Download Astronomy Astrophysics The accretion environment in Vela X-1 during a flaring period using

Astronomy
&
Astrophysics
A&A 563, A70 (2014)
DOI: 10.1051/0004-6361/201322404
c ESO 2014
The accretion environment in Vela X-1 during a flaring period
using XMM-Newton
S. Martínez-Núñez1 , J. M. Torrejón1,2 , M. Kühnel3 , P. Kretschmar4 , M. Stuhlinger4 , J. J. Rodes-Roca1,2 , F. Fürst5 ,
I. Kreykenbohm3 , A. Martin-Carrillo6 , A. M. T. Pollock4 , and J. Wilms3
1
2
3
4
5
6
Instituto Universitario de Física Aplicada a las Ciencias y las Tecnologías, University of Alicante, PO Box 99, 03080 Alicante,
Spain
e-mail: [email protected]
Departamento de Física, Ingenería de Sistemas y Teoría de la Señal, University of Alicante, PO Box 99, 03080 Alicante, Spain
Dr. Karl Remeis-Observatory & ECAP, Universitat Erlangen-Nürnberg, Sternwartstr. 7, 96049 Bamberg, Germany
European Space Astronomy Centre (ESA/ESAC), Science Operations Department, Villanueva de la Cañada (Madrid), Spain
Cahill Center for Astronomy and Astrophysics, California Institute of Technology, Pasadena CA 91125, USA
University College Dublin, Belfield, Dublin 4, Ireland
Received 30 July 2013 / Accepted 26 December 2013
ABSTRACT
We present analysis of 100 ks contiguous XMM-Newton data of the prototypical wind accretor Vela X-1. The observation covered
eclipse egress between orbital phases 0.134 and 0.265, during which a giant flare took place, enabling us to study the spectral properties
+0.42
both outside and during the flare. This giant flare with a peak luminosity of 3.92−0.09
× 1037 erg s−1 allows estimates of the physical
parameters of the accreted structure with a mass of ∼1021 g. We have been able to model several contributions to the observed spectrum
with a phenomenological model formed by three absorbed power laws plus three emission lines. After analysing the variations with
orbital phase of the column density of each component, as well as those in the Fe and Ni fluorescence lines, we provide a physical
interpretation for each spectral component. Meanwhile, the first two components are two aspects of the principal accretion component
from the surface of the neutron star, the third component seems to be the X-ray light echo formed in the stellar wind of the companion.
Key words. X-rays: binaries – pulsars: individual: Vela X-1
1. Introduction
Vela X-1 is an eclipsing high-mass X-ray binary discovered in
1967 (Chodil et al. 1967) at a distance of ∼2 kpc (Nagase 1989).
It consists of the early type supergiant HD 77581 (B0.5Iab,
R ∼ 34 R ) and a neutron star of ∼1.77 M (Rawls et al. 2011),
orbiting its companion with a period of 8.964 days. The distance
at periastron is only 0.6 stellar radii (a ∼ 53 R ; van Kerkwijk
et al. 1995) from the stellar surface, well within the acceleration zone of the stellar wind. Despite the status of the Vela X-1
system as an archetypal wind accretor, the structure of the extended atmosphere of the supergiant and large scale structure of
the stellar wind are not known in detail.
The X-ray luminosity of Vela X-1 is typically a few
1036 erg s−1 , four or more orders of magnitude brighter than
the intrinsic emission expected from the supergiant primary. It
is well explained by wind accretion from the optical companion,
with an estimated mass loss rate of 2 × 10−6 M yr−1 (Watanabe
et al. 2006). Intense X-ray variability of Vela X-1, lasting from
a few minutes to several hours, is commonly thought to be associated with changes in the accretion rate and to inhomogeneities
in the stellar wind flow (Nagase et al. 1983; Haberl & White
1990). There have been reports, on the one hand, of reductions
in the flux to less than 10% of its normal emission (“off states”),
as well as of very active states, on the other, with increases in the
flux by sometimes more than a factor of 10 (Kreykenbohm et al.
1999, 2008).
X-ray pulsations were discovered by McClintock et al.
(1976) with a fluctuating pulse period of ∼283 s. This random
spin up and down does not show any apparent connection to the
orbital phase of the neutron star around its supergiant companion
(Nagase et al. 1984), and it is believed to be caused by the transfer of angular momentum by wind accretion onto the neutron
star.
The X-ray spectrum has been described either by a power
law with a high energy cut-off beyond 15−30 keV (Nagase et al.
1986; Kreykenbohm et al. 1999) or by a negative positive exponential model (NPEX; Orlandini et al. 1998; Kreykenbohm
et al. 2002; Odaka et al. 2013). Below 3 keV, a soft excess is
observed, which was modelled by thermal Bremsstrahlung with
kT ∼ 0.5 keV by Haberl (1994) using ROSAT data. Moreover,
it was found that the temperature of the soft component depends
on the contribution of the highly absorbed hard power-law component from the neutron star. According to Hickox et al. (2004),
the observed soft excess is due to emission by photoionized or
collisionally heated diffuse gas or thermal emission from the surface of the neutron star.
Above 20 keV, cyclotron resonance scattering features
(CRSFs) between 25 and 32 keV (Makishima et al. 1992;
Choi et al. 1996; Kreykenbohm et al. 2002) and at ∼55 keV
(Kendziorra et al. 1992; Orlandini et al. 1998; La Barbera et al.
2003; Attié et al. 2004) have been reported, although the interpretation of the 25 keV feature is still sometimes debated
(Orlandini 2006).
Article published by EDP Sciences
A70, page 1 of 13
A&A 563, A70 (2014)
Observations of the system during the eclipse of the X-ray
pulsar with Tenma, ASCA, and Chandra have revealed various
fluorescent lines in addition to highly ionized lines and radiative recombination continua that imply the existence of optically
thick and clumped matter in addition to warm ionized plasma
(Sato et al. 1986; Nagase et al. 1994; Sako et al. 1999; Schulz
et al. 2002).
The structure of the wind of the system was investigated
above 20 keV by Fürst et al. (2010), who suggests that a mixture of a clumpy wind, shocks, and turbulence could explain the
measured mass distribution. These authors estimate a mass of
the wind clumps of the order of mcl = 1021 g for the giant flares.
Furthermore, in this work a log-normal distribution of the brightness was found along the orbit.
Odaka et al. (2013) report on Suzaku observations of Vela
X-1 covering the orbital phase interval 0.20–0.39 for phase zero
at T 90 . They found strong variability with LX (0.1–100 keV)
ranging from ∼0.5 to ∼9.5 × 1036 erg s−1 and variability time
scales of 1–10 ks. Assuming a clumpy wind as the source of
the variations, the observed variability indicated clump radii of
(2−20) × 1010 cm for a relative wind velocity of 400 km s−1 . For
a bright (LX = 1037 erg s−1 ), short (1 ks) flare this would lead to
a clump mass of 4 × 1019 g and an overdensity of two orders of
magnitude compared to a smooth wind model, while longer, less
intense flares would be explained naturally by density variations
of a factor of a few.
To study the structure of the extended atmosphere and the
wind in the Vela X-1 system, a 123 ks XMM-Newton (Jansen
et al. 2001) observation was performed immediately following
eclipse egress. The main goal was a precise measurement of the
evolution of the absorption column in this part of the orbit. From
historical observations (Sato et al. 1986; Lewis et al. 1992), one
expects a steep decrease in the absorption measure NH in this
part of the orbit as the neutron star emerges from behind the
companion’s extended atmosphere.
The remainder of the paper is structured as follows. In
Sect. 2, the data and the software used are described. Section 3
describes the light curves of the observation. Spectral analysis
results are presented in Sect. 4. Finally a discussion is given in
Sect. 5 and a summary and conclusions in Sect. 6.
2. Observations and data reduction
XMM-Newton observed Vela X-1 on 2006 May 25−26
(MJD 53 880.452 to MJD 53 881.876) using the European
Photon Imaging Camera (EPIC) and the Reflection Grating
Spectrometers (RGS) during all of revolution 1183 of the satellite. For this observation (ObsID 0406430201), the EPIC-MOS1
(EPIC-Metal Oxide Semi-conductor) camera was disabled to
maximize the available telemetry rate for the EPIC-pn CCD
camera, which was set up in timing mode with the thin filter.
In this mode, the camera provides timing information for each
event with a resolution of up to 1.5 ms, as well as the standard
energy resolution of about E/ΔE = 40.
Avoiding times close to the Earth’s radiation belts, we
used data between MJD 53 880.613 and MJD 53 881.768 (about
100 ks), corresponding to the orbital phase interval 0.132 to
0.270 according to the ephemeris of Kreykenbohm et al. (2008)
(see Table 1), where phase zero is defined as T 90 . When comparing our findings with other results, one should note that sometimes T 90 and sometimes the mid eclipse time T ecl are used in
the literature to define zero phase, where T ecl is later than T 90
by 0.226 d (Kreykenbohm et al. 2008). Using T ecl as reference
A70, page 2 of 13
Table 1. Ephemeris data used for timing calculations taken from
(1) Bildsten et al. (1997) and (2) Kreykenbohm et al. (2008).
T 90
Porb
a sin i
e
ω
52 974.001 ± 0.012 MJD
8.96436 ± 0.00003 d
113.89 lt-s
0.0898 ± 0.0012
152.6 ± 0.9◦
(2)
(2)
(1)
(1)
(2)
would shift the orbital phase of our data set by 0.0252 to earlier
phases.
In this paper we focus on the EPIC-pn data and its evolution
along the orbital motion of the neutron star around the supergiant
companion. The RGS data will be considered in later work. The
data reduction was performed using SASv11.0 and CCFs as of
2012 May starting from ODF level running rgsproc and epproc.
For EPIC-pn, we applied a rate-dependent CTI correction, and
to the time column a barycentric correction, as well as a correction for the Vela X-1 binary system using the ephemeris as of
Bildsten et al. (1997) and Kreykenbohm et al. (2008).
2.1. Spectral extraction
Due to the high count rate (>2500 cts/s) during the flare that
occurred as well as to the hard spectral emission characteristics of the source, we had to cope with two instrumental effects
before extracting spectra: pile-up and the improvement of the
timing mode charge transfer inefficiency due the high amount
of shifted charge (rate-dependent CTI). The first can cause a
change in spectral parameters, e.g. a hardening of the continuum
with increasing levels of pile-up, the latter a general gain/offset
shift of the complete spectrum. We mitigated the effects of pileup by considering the pattern distribution (epatplot), the maximum count rate (<800 cts/s)1 , and the stability of the spectral
parameters dependent on the columns used to extract the spectra. We found that the three columns covering the central part
of the point spread function (PSF) are significantly affected by
pile-up, and from spectral stability analysis we got indications
that their adjacent columns might still show slight and close to
insignificant pile-up effects.
Restricting the analysis to the flaring period, we extracted
spectra taken from the full PSF and from regions ignoring 3, 5,
7, 9, and 11 columns of the PSF centre (named PSF-# afterward)
and compared their modelling of the instrumental gold edge and
the relative position of the iron line, as well as general slope of
the continuum.
Without correction we found a shift of the gold edge of 93 eV
keeping the PSF centre, and we needed to exclude the central 11
columns (PSF-11) to be consistent with no shift of the gold edge.
The Fe line-energy shift continuously decreased with increasing
number of ignored columns by 132 eV from full PSF to PSF-11.
The SAS-task epfast is designed to correct for these ratedependent CTI effects seen in EPIC-pn fast modes (but not for
pile-up effects). Applying the epfast correction (using the corresponding public CCF-file EPN_CTI_0023.CCF), we still found
shifts at the gold edge of 13 eV using PSF-7, 12 eV using PSF-9,
and 7 eV using PSF-11, which is consistent with the systematic
uncertainties of the correction2. However, the Fe line energy was
1
According to the XMM-Newton Users Handbook.
XMM_CCF_REL_256 on
http://xmm2.esac.esa.int/external/xmm_sw_cal/calib/
rel_notes/index.shtml
2
0.4
0.2
0
normalized counts s−1 keV−1
S. Martínez-Núñez et al.: The accretion environment in Vela X-1 during a flaring periodusing XMM-Newton
0.5
Fig. 1. Orbit sketch of Vela X-1. The XMM-Newton 0.5−12 keV observed count rate is shown in the red curve and compared with the average light curve of the source over years of 1.5−12 keV RXTE/ASM
observations (green curve).
consistent for all spectra of PSF-5 to PSF-11, but at an energy of
about 40 eV higher compared to the uncorrected PSF-11 value.
We calculated the spectral shifts that depends on the amount
of charge per column for the flare period, and with support of the
XMM-Newton SOC we added the results to the calibration of the
rate dependent CTI, creating a non-public EPN_CTI_0024.CCF.
Using this special CCF, no more shifts at the gold edge were seen
for all spectra of PSF-5 to PSF-11, and the Fe line energy was
consistent with the uncorrected PSF-11 value. Also, the continuum slopes were consistent within errors, with PSF-5 slightly
harder than PSF-7 to PSF-11. Finally we decided to exclude the
innermost seven columns from the spectral extraction and calculated a corresponding arf response file. A comparison of the
full PSF spectrum and ARF with our PSF-7 spectrum and ARF
for the pre-flare period yielded a systematic error of about 3%
for the PSF-7 flux normalizations. For consistency, we use the
PSF-7 region for all epochs of our data set.
2.2. Background generation
Looking into the columns closest to the timing mode window
border, it was obvious that the observation does not provide a
local background area, since a) the light curve extracted from
these columns follows the overall light curve of the source and
b) its spectrum differs significantly from what is expected to
be an empty sky emission, resembling the source spectrum. A
background extracted from an earlier timing mode observation
of Vela X-1 (Obs.ID 0111030101) showed a similar pattern but
one restricted to higher energies. We therefore searched a timing mode “empty sky” and found Obs.ID 0406620201, where
no source is visible in the timing mode window. Excluding again
the innermost seven columns (identical PSF-7 extraction region),
we extracted a spectrum of Obs.ID 0406620201 and used it as
an empty sky background. Since the local background of the earlier Vela X-1 observation (Obs.ID 0111030101) shows a lower
background level than this “empty sky”, we merged these two
backgrounds for energies below ∼2.5 keV, scaled according to
1
2
Energy (keV)
5
10
Fig. 2. Different timing mode backgrounds. Black: closed filter;
cyan: this Vela X-1 (Obs. ID 0406430201, Cols. 3−12); navy
blue: Vela X-1 (Obs. ID 0111030101, Cols. 3−12); red: empty sky
(Obs. ID 0406620201 Cols. 16−33, 41−58); green: final merged
background.
exposure times and surfaces, to address the differences of absorption columns of about factor 100−1000.
3. Energy-resolved light curves and flux ratios
During the course of the observation, Vela X-1 displays strong
variations in flux, both overall and in the relative flux in different energy bands. These variations are driven by a mix of the
intrinsic variability of the X-ray source and the modulation of
the X-ray flux by absorption and scattering within the system.
These variations happen on all time scales studied, from tens of
kilo-seconds down to fractions of the pulse period. We have not
undertaken a detailed analysis of the timing behaviour on time
scales below the pulse period. A corresponding study is under
way and will be published separately. In this work a mostly stable pulse period of 283.389 ± 0.004 s is found. This pulse period
was measured using a phase-connected technique over 10 ks intervals. The mean photon arrival times for each interval were
then corrected by an empirical model and fit using a polynomial
function.
To describe the overall evolution, we have generated light
curves in several energy bands using the approach and corrections as described in Sect. 2. The bands used are 0.6–1, 1–3,
3–6, 6–8, and, 8–10 keV, respectively. Figure 3 shows the light
curves in the different bands for comparison. The ratios between
the count rates in the first four bands and the 8–10 keV band,
which is least affected by the strong absorption variations, are
displayed in Fig. 4.
Based on the light curves and their relative rates, several distinct intervals can be defined on the time scale of
tens of kilo-seconds, as explained in the following, where we
used T obs to denote the time since the start of the observation
(MJD 53 880.61305). During the first ≈10 ks the total flux varies
moderately, but as Fig. 4 shows, there are opposing trends in
the flux ratios below and above 3 keV. After this interval, the
ratio 6–8/8–10 keV remains nearly constant despite large flux
variations, indicative of a stable spectral shape for the unabsorbed source. A broad peak at all energies is evident between
T obs ≈ 10 ks and T obs ≈ 40 ks, which is most pronounced in the
1–3 keV band.
A70, page 3 of 13
A&A 563, A70 (2014)
Orbital phase
0.14
0.16
0.18
0.20
0.22
0.24
0.26
8.0 - 10.0 keV
100.00
6.0 - 8.0 keV
3.0 - 6.0 keV
1.0 - 3.0 keV
0.6 - 1.0 keV
Countrate (counts/s)
10.00
1.00
0.10
0.01
0
20
40
60
80
100
ks since start
Fig. 3. Light curve observed during the XMM-Newton observation with the EPIC-PN camera with a binning time of 141s (∼half a pulse period).
The bands shown are 0.6–1, 1–3, 3–6, 6–8 and 8–10 keV. The count rate is plotted on a logarithmic scale to emphasize the variability in the
different energy bands.
From T obs ≈ 40 ks to T obs ≈ 60 ks the source flux starts
to rise in all energy bands. On top of the general rise, there is
flaring visible on time scales of ks between T obs ≈ 45 ks and
T obs ≈ 55 ks, again most pronounced in the 1–3 keV band.
Around T obs ≈ 60 ks, the total source flux rises rapidly, more
than an order of magnitude within a few kilo-seconds. At energies ≥6 keV, the flux actually changes only by a factor of a few,
but the contribution ≤3 keV rises by several orders of magnitude,
while the flux ratio 6–8 keV/8–10 keV remains stable.
Around the peak of the flare, from T obs ≈ 62 ks to T obs ≈
70 ks, the overall flux is strongly variable with a series of
spikes and dips. These structures are highly correlated across
all bands, but variations in the flux ratios are mainly visible in
the 0.6−1 and 1−3 keV bands, such that higher global flux is
mainly driven by higher contributions from the low energies.
From T obs ≈ 70 ks to T obs ≈ 78 ks the flare decays with the
overall flux diminishing by an order of magnitude, similar in all
bands. Around T obs ≈ 80 ks, a smaller, but still significant flare
is visible, where again any spectral variation is mainly visible at
lower energies.
Beyond T obs ≈ 82 ks, up to the end of the observation, the
source seems to settle. While the flux still varies randomly in all
bands, albeit less than during earlier phases, the ratios between
different bands do not show a pronounced evolution.
A70, page 4 of 13
While the short term variability appears to be mostly dominated by the intrinsic flux variations, it is important to note that
the strong absorption in the circumstellar wind has a marked effect on the overall observed count rates. Moreover, the dramatic
change at the onset of the flare seems to be driven by a significant
variation in the absorbing and reprocessing material between the
source and the observer.
4. Orbital phase-resolved spectroscopy
To characterize the evolution of the stellar wind inhomogeneities
surrounding the neutron star, a study of the flux variations in
combination with the changes in the absorption column is carried out, under the hypothesis that these variations are proportional to the local Ṁ. Therefore we have divided our data in
88 EPIC-pn spectra of 1.1 ks exposure time each. The spectral
time resolution was chosen as a compromise between the statistics of the data and the short term variations observed in the light
curve analysis. Some of these spectra are shown in Fig. 5 at different relevant periods.
Since a fully physical model of the X-ray production mechanism in accreting X-ray pulsars does not yet exist owing to
the complexity of the physical problem (significant progress
has, however, been made, see e.g. Becker & Wolff 2007), an
S. Martínez-Núñez et al.: The accretion environment in Vela X-1 during a flaring periodusing XMM-Newton
0.18
0.20
0.22
0.24
0.26
1000
Orbital phase
Ratio 0.6-1 / 8-10 keV
1.00
100
0.10
10
0.01
0
20
0.14
0.16
40
ks since start
0.20
0.18
1
100
80
0.22
0.24
0.26
1000
Orbital phase
10
Ratio 1-3 / 8-10 keV
60
100
1
10
0
20
40
Total countrate 0.6-10 keV
0.16
60
Total countrate 0.6-10 keV
0.14
1
100
80
ks since start
0.18
0.20
0.22
0.24
0.26
1000
Ratio 3-6 / 8-10 keV
Orbital phase
10
100
10
1
0
20
0.14
0.16
40
60
ks since start
0.20
0.18
1
100
80
0.22
0.24
0.26
1000
Ratio 6-8 / 8-10 keV
Orbital phase
100
10
0
20
40
Total countrate 0.6-10 keV
0.16
60
80
Total countrate 0.6-10 keV
0.14
1
100
ks since start
Fig. 4. Ratios between count rates in different energy bands to that of the 8–10 keV band. From top to bottom: 0.6–1 keV, 1–3 keV, 3–6 keV, and
6–8 keV. The time binning is 1128 s (∼4 pulse periods), matching the spectra used in the time-resolved spectral analysis. In the background, the
total flux evolution is plotted for comparison with the light curves. Circles mark the data points corresponding to the spectra shown in Fig. 5.
A70, page 5 of 13
A&A 563, A70 (2014)
Fig. 5. Representative sample of the 88 EPIC-pn spectra at different relevant times marked in the inset light curve plot. The different components of
the model are shown: component 1 (blue), component 2 (green) component 3 (red) and fluorescence lines modelled as Gaussian functions (grey).
empirical model has to be used. After trying different models,
we have found that there is only one continuum model that could
fit all the spectra significantly well throughout the whole observation including the flare. It consists of three absorbed powerlaw components with a unique photon index but different normalization factors for each component. Furthermore, the spectra
of the source clearly show the presence of three emission lines
corresponding to Fe Kα , Fe Kβ , and Ni Kα fluorescence lines.
These lines have been treated as being absorbed by the highest
absorbed component and modelled by a Gaussian function even
if their shapes are not perfectly Gaussian. A full description of
the model is given in Eq. (1).
A similar continuum model has been used previously to
characterize the spectra of several high-mass X-ray binary systems, such as 4U 1700-37 (van der Meer et al. 2005), Cen X-3
(Ebisawa et al. 1996), and 4U 1538-52 (Rodes-Roca et al. 2011).
Our model differs from those models in two aspects: (a) we used
A70, page 6 of 13
a more complex ISM absorption model (tbnew in Eq. (1)) and (b)
we applied a non-relativistic Compton scattering correction to
the first and second spectral components because of their large
measured hydrogen column densities (cabs in Eq. (1)).
F(E) =
2 tbnewi × cabsi × Normi × E −Γ
i=1
+ tbnew3 × Norm3 × E −Γ
3 +
tbnew1 × cabs1 × Gaussian j .
(1)
j=1
In Eq. (1) the summation index i refers to the three model components, and the index j refers to the three emission lines components. Normi are the normalization factors of each power law
(photons keV−1 cm−1 s−1 at 1 keV), and Γ is the power law
S. Martínez-Núñez et al.: The accretion environment in Vela X-1 during a flaring periodusing XMM-Newton
photon index. These components are further modified by two
effects both arising in the material between the X-ray source and
the observer with density column NH , on one hand, the photoelectric absorption. For this purpose we use the model tbnew, an
updated version of the Tübingen-Boulder ISM absorption model
(Wilms et al. 2000). The absorption cross sections are adapted
from Verner et al. (1996), and the abundances are set to those
of Wilms et al. (2000). On the other hand, the non-relativistic
Compton scattering effect is taken into account by using a cabs
model, which is described by
cabs(E) = exp(−NH σT (E)),
(2)
where σT (E) is the Thomson cross section, and NH is the equivalent hydrogen column that it is linked to its corresponding absorption law.
In time- or phase-resolved spectroscopy, each spectrum is
usually fitted individually by minimizing the statistic, while
here parameters are varied, and a minimum is sought across
all 88 spectra simultaneously using the Interactive Spectral
Interpretation System (ISIS) software (Houck & Denicola 2000).
This novel approach allows finding a common continuum, as
well as obtaining the best possible fit statistics and the best constraints of the relevant parameters. Figure 5 show the spectra
and fitted models for some representative time bins throughout
the observation. The inspection of this plot already reveals that
the most dramatic changes are produced in components 2 and 3.
In a first analysis we left the power law index Γ free to vary
among all 88 spectra. We found no significant differences among
the best-fit values of this parameter for any of the 88 spectra.
Consequently, in order to increase the precision of the other parameters, we assume a common value for Γ, i.e. a stable continuum slope for all 88 spectra. The best-fit value for this index is
Γ = 1.595 ± 0.010 (χ2 = 14011 with 9765 d.o.f.). The absorption
column NH(3) of the third component was found to be constant
within errors. We therefore kept it fixed to the average value of
(0.75 ± 0.03) × 1022 cm−2 . This value is in excellent agreement
with the ISM absorption towards GP Vel, the optical counterpart of Vela X-1, derived from optical data. Finally, to pin down
the fluorescence lines, we have constrained the Fe Kβ, as well
as the Ni Kα energies, with respect to the Fe Kα centroid energy, taking theoretical calculations into account (Kallman et al.
2004). The energy of the Fe Kα line (left free) turns out to be
6.435 ± 0.001 keV. The energy of the Fe Kβ line is shifted by
0.65 keV, while its flux is considered to be 13% of that of Kα.
These energies are compatible with K shell fluorescence from
ionized Fe up to a maximum of Fexviii (Kallman et al. (2004),
Fig. 3). In the same way, the Ni Kα centroid energy was shifted
by 1.2 keV with respect to that of Fe Kα. These ratios were kept
fixed during the fitting process.
To estimate the possible dependence among the relevant parameters during the fits, we have computed confidence level contour maps for all 88 spectra. In Fig. 6 we show contour maps for
the same selected spectra as in Fig. 5. As can be seen, the contours do not betray a systematic dependence, and therefore, the
obtained parameters and their errors are considered reliable.
4.1. Evolution of the parameters
In Fig. 7 we show the evolution of the relevant parameters of
our model throughout time and orbital phase. Panel a shows the
unabsorbed flux evolution of each component and the overall
unabsorbed flux (note the logarithmic ordinate scale). A clear
progressive increase in the fluxes of components 2 and 3 can be
seen, and meanwhile the flux of component 1 decreases slightly
from the beginning to the end of the observation. On top of this
we clearly observe two flares at φ ≈ 0.17 and 0.22 (a giant flare),
respectively. A vertical grey line marks the rise of the bright
flare.
In panel b we show the relative unabsorbed flux evolution
of each power law against the average flux of the first component. Several characteristics can be pointed out here. On one
hand, as stated before, the most dramatic changes are seen in
components 2 (scattered) and 3 (low energy). During the bright
flare, component 3 increases by two orders of magnitude, while
component 2 increases one order of magnitude. These two components are very well correlated throughout the whole observation (except at the maximum of the bright flare). These can be
clearly seen in panel C of Fig. 8. In contrast, component 1 (direct) increases only by a factor of 2 during the flare. On the other
hand, while components 2 and 3 show an overall increase in flux
throughout the observation, following the trend of the light curve
(top panel), component 1 shows an overall decrease in flux from
1 × 10−8 erg s−1 cm−2 to 0.2 × 10−8 erg s−1 cm−2 . After the flare,
the flux of components 1 and 2 become identical within the errors, and in fact, the flux of all three components become very
well correlated. This is clearly seen in panels B and C of Fig. 8.
In panel c of Fig. 7 we show the evolution of the column
density of the absorbing material for components 1 and 2 in
units of 1022 cm−2 . The NH(3) for component 3 remains constant
throughout the observation at ∼0.75 × 1022 cm−2 , which is entirely compatible with the ISM value (van Genderen 1981). As
can be seen, the absorption of component 1 (direct) decreases
from values of 1024 cm−2 at the beginning of the eclipse egress
down to ∼1.8 × 1023 cm−2 at quadrature. This reflects the decreasing amount of stellar wind material in the line of sight towards the neutron star as it emerges from eclipse. An enhancement is present, however, between orbital phases 0.19 and 0.23.
Remarkably, there is no dramatic change in NH(1) at the moment of the flare. In contrast, NH(2) , which stays essentially constant throughout the observation, decreases drastically exactly
at the rise of the flare. After that, it remains constant again at
∼2.5 × 1022. Since the fluxes shown before were unabsorbed and
NH(3) is found to be constant during the flare (during the whole
observation, in fact), the sudden flux increase of component 3
must be intrinsic. On the other hand, the flux increase in component 2 could be either intrinsic or due to the sudden decrease
in NH(2) .
In panels d and e of Fig. 7 we show the change in the intensities and equivalent width, respectively, of the fluorescence Fe Kα
and Ni Kα lines. As can be seen, the intensity of the Fe Kα line
follows the overall brightness of the source well, including the
sharp rise at the flares (note again the logarithmic scale in the
ordinates). This is clearly seen in panel D of Fig. 8. However,
the most striking fact is that the Fe Kα line is clearly less intense
after the flare than before. This is also clearly seen in panel D of
Fig. 8 as a hysteresis cycle. The equivalent width, in turn, which
is related with the density column of the reprocessing material,
follows the increase in continuum flux up to the flare, when it
reaches a value of ∼230 eV, and then decreases in the post flare,
staying more or less constant, at a a value of 180 eV approximately, until the end of the observation (panel E, Fig. 8). This
behaviour must arise from a depletion of neutral Fe in the circumsource matter after the flare. We recall that the energy of the
Fe Kα lines remains constant throughout the observation and
shows no sign of increasing ionization. Therefore, the depletion
in neutral Fe must reflect a real depletion of circumsource matter.
A70, page 7 of 13
A&A 563, A70 (2014)
Fig. 6. Contour maps (2σ confidence limits for
the two selected parameters) between the free
spectral parameters for the six selected spectra
shown in Fig. 5. These maps demonstrate that
while for some spectra the parameters cannot
be constrained tightly, the observed changes of
parameter values along the observation are real
and significant.
Regarding the Ni Kα line, its parameters show the same trend as
those of the Fe Kα line, albeit with larger uncertainties.
4.2. Correlations among parameters
As mentioned in the previous section, we observe different correlations among the continuum parameters, looking in detail at
panels B and C of Fig. 8, three branches are clearly identified:
– Before the first peak, all fluxes decrease but component 1
and 3 do so more rapidly than component 2.
A70, page 8 of 13
– From the first peak to the pre-flare (red to blue points) all
fluxes correlate.
– During the flare and later, the fluxes correlate again, but with
a different slope than before.
The same three branches are observed when we plot the Fe Kα
fluorescence line area against the normalization factor of the second component (see panel D of Fig. 8). Owing to the larger uncertainties, there is no clearly visible correlation when plotting
S. Martínez-Núñez et al.: The accretion environment in Vela X-1 during a flaring periodusing XMM-Newton
Fig. 7. From top to the bottom: evolution of the model parameters versus time and orbital phase. Error bars are 90% confidence level. The dashed
line in panel c) indicates the constant NH of component 3. The grey column marks the rise of the flare and the grey curve in panel a) shows the
overall unabsorbed flux of the spectral model.
A70, page 9 of 13
A&A 563, A70 (2014)
Orbital phase
0.14
0.16
0.18
0.20
0.22
Total unabsorbed flux
A
0.24
0.26
58
48
10
15
5
72
85
0
20
40
60
80
100
ks since start
25
B
3.0
2.5
Flux component 3
Flux component 1
20
C
15
10
14
12
10
8
6
4
2
0
5
0
5
1.5
0.12
0.10
0.08
0.06
0.04
0.02
0.00
-0.02
1.0
0.5
0
0
2.0
10
15
Flux component 2
2
4
6
20
0.0
8
25
0
0
0.4
5
10
15
Flux component 2
2
4
6
20
8
25
300
D
E
250
EqW Fe Kα line (eV)
Flux Fe Kα line
0.3
0.2
200
150
100
0.1
50
0.0
0
0
5
10
Flux component 2
15
0
5
10
Flux component 2
15
Fig. 8. A) 0.6−10 keV unabsorbed flux light curve with a binning time of 1.1 ks. Marked circles indicate the data points corresponding to the
spectra shown in Figs. 5 and 6. B) Flux of the first component versus that of the second component. The inset shows an enlarged view of the first
58 ks, before the onset of the giant flare. C) Flux of the third component versus that of the second component with an inset equivalent to that in B).
D) Fe Kα fluorescence line flux areas against normalization factor of second component. E) Fe Kα equivalent width against the normalization
factor of the second component. All fluxes are in units of 10−9 erg cm−2 s−1 .
A70, page 10 of 13
S. Martínez-Núñez et al.: The accretion environment in Vela X-1 during a flaring periodusing XMM-Newton
between 0.5 and 600 keV of Vela X-1 during the XMM-Newton
observation is
1
1
37
−1
L = 3.92+0.42
−0.09 × 10 erg s .
Based on hard X-ray data from INTEGRAL observations, Fürst
et al. (2010) found a log-normal distribution of flux values, with
a median absolute luminosity LX = 5.1 × 1036 erg s−1 and
multiplicative standard deviation σ ≈ 2. Even though a direct comparison with Fürst et al. (2010) is difficult because of
possible variations in the continuum, it is clear that the flare
around orbital phase 0.22 is brighter than ∼99.9% of the 3.6 Ms
INTEGRAL monitoring data, and thus certainly is one of the
rare giant flares.
D b
N
3.8
4.0
2
4.6
4.8
Luminosity ({ keV) (7 erg s )
5.0
Fig. 9. Distribution of X-ray luminosities for the peak of the flare in the
energy range 0.5–600 keV using simulated spectra.
the equivalent width of this line instead (panel E of Fig. 8). Still,
two different regions are observed. During the rise and first peak
of the flare the equivalent width reaches its highest value, with no
clear variation despite significant changes in the flux of the second component. Before and after the flare, there is an apparent
correlation between these two parameters. There is an indication
for a different correlation coefficient before and after the flare,
but the relatively large uncertainties of the equivalent width do
not allow drawing a firm conclusion.
5. Discussion
5.1. Source brightness and variability
To put the current observation into perspective, it is instructive to
compare the obtained fluxes and derived luminosities with other
observations. Such comparisons, however, depend strongly on
assumptions of the overall spectral shape, which for many historical data sets is not well constrained, leading quickly to systematic uncertainties of a few 10% or more.
For this reason, we assumed the spectrum of Vela X-1 above
10 keV to be modified by a Fermi-Dirac cutoff (Tanaka 1986),
as found by Kreykenbohm et al. (2008), among others. Since
the luminosity strongly depends on the parameters used to describe that high energy cutoff, we decided to take those parameters listed in Table 5 of Kreykenbohm et al. (2008), which result in a very soft broad-band spectrum. Choosing another set of
parameters led to a harder spectrum and an increase in the luminosity estimation at those energies, which might overestimate
the real source luminosity. The assumed parameters are there+0.5
fore Ecut = 35.6+7.5
−11.5 keV and E fold = 11.2−0.3 keV. The power
law photon index Γ is consistent with the value presented here
(see Sect. 4).
To derive proper values for the uncertainty of the estimated
luminosity, we have to take the uncertainties of the fit parameters
into account, as well as those of the assumed FD cutoff. Based on
a Monte Carlo approach, we calculated the luminosity of spectrum 56 (the peak of the flare) with varying random parameters,
corresponding to their individual asymmetric uncertainties. The
resulting simulated distribution of luminosities between 0.5 and
600 keV after 100 000 runs is shown in Fig. 9. The uncertainty
interval of the most-likely value was chosen such, that it contains
90% of the distribution. The final value of the peak luminosity
5.2. Probable origin of the spectral components
Although we have used a phenomenological model, the behaviour of the three power law components in Fig. 7 with respect
to the flare immediately suggests a further unifying idea of only
two physical sources: the neutron-star surface and the scattered
emission by electrons in the wind. Continuum components 2
and 3 suffer parameter discontinuities across the flare with significant increases in the flux relative to component 1 coinciding
with a steep decrease in component 2’s absorption such that the
post-flare fluxes of components 1 and 2 are essentially identical
during the subsequent decline of about an order of magnitude in
brightness. This suggests that components 1 and 2 are two aspects of the same component that originates at the neutron-star
surface. Absorption along the line-of-sight through the wind is
complex owing to the warm ionization state of the supergiant’s
expanding envelope but can be modelled successfully by the sum
of two conventional cold absorption laws. The balance between
these two changed across the flare because photoionization altered the ionic balance in the wind by removing electrons from
the ions responsible for absorption near 1 keV. These extra free
electrons were then responsible for the extra scattering that produced the increase in component 3.
As we have seen, NH(1) decreases from 1024 at orbital phase
φ = 0.14 to 1023 at orbital phase φ = 0.26 as the neutron star
emerges from eclipse and moves along the orbit (Fig. 7, third
panel). However, there is clearly a wind density enhancement
starting at φ ≈ 0.185, followed by an increase in the emission of
the neutron star due probably to the corresponding gradual accretion of more mass. The later, more sudden increase in emission from the neutron star was probably due to the accretion of a
dense clump in the wind, although this was accompanied by no
change of spectral index. The time scale for the rise of the flare
was of the order of one time bin or trise ∼ 103 s.
5.3. Physical parameters of wind clumps
Our observation allows us to make direct estimates of the physical properties of the possible clump responsible for the giant
flare. Considering a spherical shape of the clump and taking the
characteristic time for the accretion of the clump into account
would be trise ∼ 103 s. Assuming that the clumps are moving
with the same velocity as the bulk of the wind, they are travelling at v(r) = v∞ (1 − R∗ /r)β = 0.46v∞ = 5.1 × 107 cm/s, with
v∞ = 1100 km s−1 , r = 1.6 R∗ and β = 0.8. Therefore, the size of
the clump is of the order of
lcl ∼ trise v 5 × 1010 cm
(3)
A70, page 11 of 13
A&A 563, A70 (2014)
or ∼0.02R∗, in agreement with theoretical expectations from
massive star winds (Oskinova et al. 2011, Table 1).
Using a finer time bin, as that used in Fig. 3 (141 s, half the
NS spin period) we still see four or five data points during the
3
rise of the big flare compatible with a trise <
∼ 10 s. We can now
estimate the density of the clump from
ncl ∼
NHwind 5 × 1023
≈
≈ 1013 H atoms/cm3 = 10−11 g/cm3.
lcl
5 × 1010
(4)
On the other hand, the characteristic volume of the clump would
be Vcl ∼ lcl 3 ≈ 1032 cm3 . Therefore, the mass of the clump
responsible of the bright flare would be mcl ∼ 1021 g.
This value is two orders of magnitude greater than deduced
by Odaka et al. (2013; Eq. (13)). As we have argued, however,
the flare we are analysing here is much brighter than the one
seen by Suzaku. Our estimate agrees with the one by Fürst et al.
(2010) for clumps responsible for bright flares like the one we
have observed. We can also estimate the luminosity produced
by the accretion of such a clump onto the neutron star (NS).
Since this clump accretes in a characteristic time of trise ∼ 103 s,
the mass accretion rate will be Ṁ ∼ 1018 g/s. Now, if the mass
accretion-to-energy conversion efficiency η adopts reasonable
values [0.1−0.3], the luminosity of the flare should then be of
the order of
MNS Ṁ
LX = ηG
≈ [2 − 6] × 1037 erg/s
(5)
RNS
as observed. In conclusion the bright flare observed by
XMM-Newton is consistent with the accretion of a dense clump
whose size is otherwise the one expected for the structured stellar wind of a B0.5Ib star at r <
∼ 2R∗ .
6. Summary and conclusions
The observation during eclipse egress has revealed unexpectedly rich behaviour of the Vela X-1 system. From the light curve
analysis, we found strong variations in source flux (overall and
relative flux in different bands) on all time scales studied from
kilo-seconds down to fractions of the pulse period. At orbital
phase 0.22, a giant flare took place, which was mainly driven by
a change in the absorbing and reprocessing material between the
source and the observer. Even though the observed short-term
variability in the light curves appears to be mostly dominated by
intrinsic flux variations, the spectral shape of the system remains
rather stable along the observation, reflecting the marked effect
of the strong absorption in the circumstellar wind on the overall
observed count rates.
The observation was divided into 88 spectra of 1.1 ks exposure time each to study the spectral evolution of the system along
the orbit. A phenomenological model allows us to fit all spectra well throughout the whole observation, including the flare. It
consists of three continuum components formed by an absorbed
power-law, with the same power-law index for each component.
This continuum spectral model is complemented by three fluorescent emission lines of Fe Kα , Fe Kβ , and Ni Kα . The main
results from the spectral analysis follow:
– The power-law index remains constant along the observation
with a best value of Γ = 1.595 ± 0.010.
– The overall unabsorbed fluxes of components 2 and 3 increase during the observation, and meanwhile the overall unabsorbed flux of component 1 slightly decreases from the
beginning to the end of the observation.
A70, page 12 of 13
– Two flares are clearly observed at φ ≈ 0.17 and 0.22 (giant
flare) in the unabsorbed fluxes of the three components.
– During the rise of the giant flare, which lasted 1.1 ks, the
most dramatic changes were in components 3 and 2, with
two and one order of magnitude increases in their unabsorbed fluxes, respectively.
– The absorption column of the first component decreases
from values 1024 cm−2 at the beginning of eclipse egress up
to ∼1.8 × 1023 cm−2 at quadrature. Meanwhile, the absorption column of the second component stays constant from the
beginning of the observation until the rise of the flare, when
it decreases drastically from 1023 cm−2 to 2.5 × 1022 cm−2 .
On the other hand, the absorption column of the third component remains constant throughout of the observation with
a best value of (0.75 ± 0.03) × 1022 cm−2 , in excellent agreement with the ISM absorption towards HD 77581, the optical
counterpart of Vela X-1.
– The energy of the Fe Kα line was constant during the observation with a best value of 6.435 ± 0.001 keV, reflecting no
sign of increasing ionization, even during the giant flare.
– The Fe Kα line flux follows the overall brightness of the
source well, is clearly less intense after the flare. Moreover,
the equivalent width of the line follows the increase in continuum flux up to the flare, reaching a maximum value of
∼230 eV, and decreasing after the flare achieving a constant
value of 180 eV. This observed behaviour of the line parameters clearly indicates a depletion of the neutral Fe in the
circumsource matter after the giant flare.
– The unabsorbed fluxes of three components are correlated
very well during the whole observation except at the peak of
the flare. We could establish three different regimes of correlation among these parameters, which are also observed
between the Fe Kα flux line and the flux of the second
component.
While our results improve understanding of the stellar wind
structure of Vela X-1 and similar sources, they do not allow inferring whether the accreted dense clump of matter is caused by
clumping in the stellar wind of the massive star, independent of
the accretion process, or by density perturbations caused by the
passage of the accretor through the dense wind as found in hydrodynamical simulations such as those of Blondin et al. (1991)
and Manousakis et al. (2012). Further insights could be gained
from detailed simulations that take all these effects into account
in a self-consistent manner, but they do not exist yet.
Acknowledgements. This work was supported by the Spanish Ministerio de
Ciencia e Innovación through the projects AYA2010-15431 and AIB2010DE00054. It was partly supported by the Bundesministerium für Wirtschaft
und Technologie under Deutsches Zentrum für Luft- und Raumfahrt grant
50OR1113. This research was made possible in part by a travel grant from the
Deutscher Akademischer Austauschdienst. JJRR acknowledges the support by
the Vicerectorat d’Investigació, Desenvolupament i Innovació de la Universitat
d’Alacant under grant GRE12-35. The authors acknowledge the help of the
International Space Science Institute at Bern, Switzerland and support by the
Faculty of the European Space Astronomy Centre. The SLXfig package, developed by John E. Davis, was used to produce some of the figures within in this
paper. We thank the anonymous referee whose comments allowed us to improve
this paper.
References
Attié, D., Chernyakova, M., Kretschmar, P., & et al. 2004, in 35th COSPAR
Scientific Assembly, 2862
Becker, P. A., & Wolff, M. T. 2007, ApJ, 654, 435
Bildsten, L., Chakrabarty, D., Chiu, J., et al. 1997, ApJS, 113, 367
Blondin, J. M., Stevens, I. R., & Kallman, T. R. 1991, ApJ, 371, 684
S. Martínez-Núñez et al.: The accretion environment in Vela X-1 during a flaring periodusing XMM-Newton
Chodil, G., Mark, H., Rodrigues, R., Seward, F. D., & Swift, C. D. 1967, ApJ,
150, 57
Choi, C. S., Dotani, T., Day, C. S. R., & Nagase, F. 1996, ApJ, 471, 447
Ebisawa, K., Day, C. S. R., Kallman, T. R., et al. 1996, PASJ, 48, 425
Fürst, F., Kreykenbohm, I., Pottschmidt, K., et al. 2010, A&A, 519, A37
Haberl, F. 1994, A&A, 288, 791
Haberl, F., & White, N. E. 1990, ApJ, 361, 225
Hickox, R. C., Narayan, R., & Kallman, T. R. 2004, ApJ, 614, 881
Houck, J. C., & Denicola, L. A. 2000, in Astronomical Data Analysis Software
and Systems IX, eds. N. Manset, C. Veillet, & D. Crabtree, ASP Conf. Ser.,
216, 591
Jansen, F., Lumb, D., Altieri, B., et al. 2001, A&A, 365, L1
Kallman, T. R., Palmeri, P., Bautista, M. A., Mendoza, C., & Krolik, J. H. 2004,
ApJS, 155, 675
Kendziorra, E., Mony, B., & Kretschmar, P. 1992, in Frontiers Science Series,
Proc. Yamada Conf. XXVIII, eds. Y. Tanaka, & K. Koyama (Tokyo, Japan:
Universal Academy Press), 51
Kreykenbohm, I., Kretschmar, P., Wilms, J., et al. 1999, A&A, 341, 141
Kreykenbohm, I., Coburn, W., Wilms, J., et al. 2002, A&A, 395, 129
Kreykenbohm, I., Wilms, J., Kretschmar, P., et al. 2008, A&A, 492, 511
La Barbera, A., Santangelo, A., Orlandini, M., & Segreto, A. 2003, A&A, 400,
993
Lewis, W., Rappaport, S., Levine, A., & Nagase, F. 1992, ApJ, 389, 665
Makishima, K., Mihara, T., Nagase, F., & Murakami, T. 1992, in Frontiers
Science Series, Proc. Yamada Conf. XXVIII, eds. Y. Tanaka, & K. Koyama
(Tokyo, Japan: Universal Acadamy Press), 23
Manousakis, A., Walter, R., & Blondin, J. M. 2012, A&A, 547, A20
McClintock, J. E., Rappaport, S., Joss, P. C., et al. 1976, ApJ, 206, L99
Nagase, F. 1989, PASJ, 41, 1
Nagase, F., Hayakawa, S., Makino, F., Sato, N., & Makishima, K. 1983, PASJ,
35, 47
Nagase, F., Hayakawa, S., Kunieda, H., et al. 1984, ApJ, 280, 259
Nagase, F., Hayakawa, S., Sato, N., Masai, K., & Inoue, H. 1986, PASJ, 38,
547
Nagase, F., Zylstra, G., Sonobe, T., et al. 1994, ApJ, 436, L1
Odaka, H., Khangulyan, D., Tanaka, Y. T., et al. 2013, ApJ, 767, 70
Orlandini, M. 2006, in Spectra & Timing of Compact X-ray Binaries, Adv. Space
Res., 38, 2742
Orlandini, M., dal Fiume, D., Frontera, F., et al. 1998, A&A, 332, 121
Oskinova, L., Hamann, W.-R., Ignace, R., & Feldmeier, A. 2011, Bull. Soc. Roy.
Sci. Liège, 80, 54
Rawls, M. L., Orosz, J. A., McClintock, J. E., et al. 2011, ApJ, 730, 25
Rodes-Roca, J. J., Page, K. L., Torrejón, J. M., Osborne, J. P., & Bernabéu, G.
2011, A&A, 526, A64
Sako, M., Liedahl, D. A., Kahn, S. M., & Paerels, F. 1999, ApJ, 525, 921
Sato, N., Hayakawa, S., Nagase, F., et al. 1986, PASJ, 38, 731
Schulz, N. S., Canizares, C. R., Lee, J. C., & Sako, M. 2002, ApJ, 564, L21
Tanaka, Y. 1986, in Radiation Hydrodynamics in Stars and Compact Objects,
Proc. IAU Colloq. 89, eds. D. Mihalas, & K.-H. A. Winkler (Springer-Verlag),
Lect. Notes Phys., 255, 198
van der Meer, A., Kaper, L., di Salvo, T., et al. 2005, A&A, 432, 999
van Genderen, A. M. 1981, A&A, 96, 82
van Kerkwijk, M. H., van Paradijs, J., Zuiderwijk, E. J., et al. 1995, A&A, 303,
483
Verner, D. A., Ferland, G. J., Korista, K. T., & Yakovlev, D. G. 1996, ApJ, 465,
487
Watanabe, S., Sako, M., Ishida, M., et al. 2006, ApJ, 651, 421
Wilms, J., Allen, A., & McCray, R. 2000, ApJ, 542, 914
A70, page 13 of 13